University of Groningen
Chasing discs around O-type (proto)stars
Maud, L. T.; Cesaroni, R.; Kumar, M. S. N.; van der Tak, F. F. S.; Allen, V.; Hoare, M. G.;
Klaassen, P. D.; Harsono, D.; Hogerheijde, M. R.; Sanchez-Monge, A.
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Astronomy & astrophysics DOI:
10.1051/0004-6361/201833908
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Publication date: 2018
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Maud, L. T., Cesaroni, R., Kumar, M. S. N., van der Tak, F. F. S., Allen, V., Hoare, M. G., Klaassen, P. D., Harsono, D., Hogerheijde, M. R., Sanchez-Monge, A., Schilke, P., Ahmadi, A., Beltran, M. T., Beuther, H., Csengeri, T., Etoka, S., Fuller, G., Galvan-Madrid, R., Goddi, C., ... Zinnecker, H. (2018). Chasing discs around O-type (proto)stars. Astronomy & astrophysics, 620, [31].
https://doi.org/10.1051/0004-6361/201833908
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Astronomy
&
Astrophysics
https://doi.org/10.1051/0004-6361/201833908© ESO 2018
Chasing discs around O-type (proto)stars
ALMA evidence for an SiO disc and disc wind from G17.64+0.16
?
L. T. Maud
1, R. Cesaroni
2, M. S. N. Kumar
3,4, F. F. S. van der Tak
5,6, V. Allen
5,6, M. G. Hoare
7,
P. D. Klaassen
8, D. Harsono
1, M. R. Hogerheijde
1,9, Á. Sánchez-Monge
10, P. Schilke
10, A. Ahmadi
11,
M. T. Beltrán
2, H. Beuther
11, T. Csengeri
12, S. Etoka
13, G. Fuller
13,14, R. Galván-Madrid
15, C. Goddi
1,16,
Th. Henning
11, K. G. Johnston
7, R. Kuiper
17, S. Lumsden
7, L. Moscadelli
2, J. C. Mottram
11, T. Peters
18,
V. M. Rivilla
2, L. Testi
19,2, S. Vig
20, W. J. de Wit
21, and H. Zinnecker
22,23(Affiliations can be found after the references) Received 18 July 2018 / Accepted 13 September 2018
ABSTRACT We present high angular resolution (∼0.200
) continuum and molecular emission line Atacama Large Millimeter/sub-millimeter Array (ALMA) observations of G17.64+0.16 in Band 6 (220−230 GHz) taken as part of a campaign in search of circumstellar discs around (proto)-O-stars. At a resolution of ∼400 au the main continuum core is essentially unresolved and isolated from other strong and compact emission peaks. We detect SiO (5–4) emission that is marginally resolved and elongated in a direction perpendicular to the large-scale outflow seen in the13CO (2−1) line using the main ALMA array in conjunction with the Atacama Compact Array (ACA).
Morphologically, the SiO appears to represent a disc-like structure. Using parametric models we show that the position-velocity profile of the SiO is consistent with the Keplerian rotation of a disc around an object between 10 and 30 M in mass, only if there is also
radial expansion from a separate structure. The radial motion component can be interpreted as a disc wind from the disc surface. Models with a central stellar object mass between 20 and 30 M are the most consistent with the stellar luminosity (1 × 105L ) and
indicative of an O-type star. The H30α millimetre recombination line (231.9 GHz) is also detected, but spatially unresolved, and is indicative of a very compact, hot, ionised region co-spatial with the dust continuum core. The broad line-width of the H30α emission (full-width-half-maximum = 81.9 km s−1) is not dominated by pressure-broadening but is consistent with underlying bulk motions.
These velocities match those required for shocks to release silicon from dust grains into the gas phase. CH3CN and CH3OH thermal
emission also shows two arc shaped plumes that curve away from the disc plane. Their coincidence with OH maser emission suggests that they could trace the inner working surfaces of a wide-angle wind driven by G17.64 which impacts the diffuse remnant natal cloud before being redirected into the large-scale outflow direction. Accounting for all observables, we suggest that G17.64 is consistent with a O-type young stellar object in the final stages of protostellar assembly, driving a wind, but that has not yet developed into a compact HIIregion. The existance and detection of the disc in G17.64 is likely related to its isolated and possibly more evolved nature, traits which may underpin discs in similar sources.
Key words. stars: formation – stars: protostars – stars: massive – stars: winds, outflows – stars: pre-main sequence – submillimeter: stars
1. Introduction
The formation scenario for the most massive stars (>8 M ) still
remains uncertain. From a theoretical stand point models gen-erally invoke a scaled-up solar-mass model, with a disc, or disc-like structure, accretion streams and axial jets or outflows (e.g. Krumholz et al. 2009; Peters et al. 2010a; Kuiper et al. 2010,2011;Klassen et al. 2016;Rosen et al. 2016;Tanaka et al. 2017; Kuiper & Hosokawa 2018) as otherwise massive young stellar objects (YSOs) could not accrete in the face of intense stellar winds and radiation pressure (e.g. Wolfire & Cassinelli 1986;Nakano 1989; Jijina & Adams 1996). Observations indi-cate that jets and outflows are associated with massive YSOs (e.g.Beuther et al. 2002;Maud et al. 2015b;Purser et al. 2016;
Bally 2016), however observations of the putative discs in the mm-regime are scarce, as somewhat expected considering that
?
The reduced datacubes are only available at the CDS via anonymous ftp tocdsarc.u-strasbg.fr(130.79.128.5) or via
http://cdsarc.u-strasbg.fr/viz-bin/qcat?J/A+A/620/A31
pre-ALMA, mm-interferometers did not typically resolve targets below ∼1000 au, except in a few studies (e.g.Jiménez-Serra et al. 2012;Wang et al. 2012;Maud & Hoare 2013;Beuther et al. 2013;
Cesaroni et al. 2014;Hunter et al. 2014).Beltrán & de Wit(2016) provide a comprehensive review of discs around luminous YSOs, in which B-type sources are indicated to have characteristics of scaled-up solar-type systems. Based on previous studies the authors reported that young O-type stellar sources (L > 105L )
show large-scale rotation in ∼10 000 au, potentially transient, toroidal “pseudo-disc” structures (see alsoFuruya et al. 2008), although relatively recent ALMA observations byJohnston et al.
(2015) did indicate a Keplerian signature in a 2000 au radius structure around a ∼25 M O-type source. Keplerian-like
sig-natures are also found around B-type binary systems, due to a rotating circumbinary disc (e.g. Sánchez-Monge et al. 2013a,
2014; Beltrán et al. 2016), while O-type star forming regions show much more complex structures at millimetre wavelengths with anisotropic accretion and spiral-like feeding filaments from large parsec scales to small ∼1000 au scales (e.g.Liu et al. 2015;
to the streams seen in the aforementioned models. Critically, the latter examples do not show evidence for Keplerian-like discs, at least down to the 500 au scales that these authors probe.
Any massive stellar source from ∼8 M can burn hot enough
to completely ionise the surrounding molecular material to form an HII region (Wood & Churchwell 1989; Churchwell 1990;
Kurtz 2005) and destroy any complex chemical tracers that are traditionally used to understand the kinematics of their natal environments and their accretion discs. Initial arguments by
Walmsley (1995) argued that high accretion rates will cause a very dense HIIregion that is optically thick to radio emission and thus the HIIwould not be seen. However, for heavily accret-ing massive YSOs the onset of the HIIregion could be delayed via stellar bloating (Palla & Stahler 1992;Hosokawa & Omukai 2009; Hosokawa et al. 2010; Kuiper & Yorke 2013) where the effective temperature of the star is much cooler than it would be considering a main sequence star of the same luminosity. In these models, a halt or considerable reduction in accretion ( 10−3M
yr−1), or growth beyond ∼30−40 M , will result in
the YSO contracting to a “main-sequence” configuration, heat-ing significantly, and beheat-ing able to create an HII region. The fine details of this transition are still somewhat unclear since they depend on the assumed accretion law and the initial condi-tions chosen for the stellar evolution calculacondi-tions (Haemmerlé & Peters 2016). Observations in the infra-red (IR) have been made in search of cool stellar atmospheres that point to bloated stars, however these studies remain inconclusive (Linz et al. 2009;
Testi et al. 2010). Alternative scenarios are that HIIregions can be gravitationally trapped at very early ionisation stages (see
Keto 2003,2007), or flicker due to chaotic shielding of the ionis-ing radiation by an accretion flow, leadionis-ing to a non-monotonous expansion (Peters et al. 2010b). Hyper Compact (HC) HIIregions (<0.03 pc) are thought to be the earliest ionisation stage and therefore could relate to the halt of accretion, be a marker of a transition phase.
Radio observations provide some evidence of the transi-tion stages. Insofar, three high-mass YSOs, S106 IR (Hoare et al. 1994), S140 IRS1 (Hoare 2006;Maud & Hoare 2013) and G298.2620+00.7394 (as a candidate fromPurser et al. 2016) are potentially at a stage contracting down from bloated stars and now driving ionised disc winds. These objects have compact, but elongated, radio continuum emission inconsistent with a jet scenario. The compact radio emission for these sources can be understood as an equatorially driven radiative disc wind (e.g.
Drew et al. 1998), where the intense radiation from the now con-tracted YSO is ionising and ablating the disc surface and driving a wind in the equatorial plane (see alsoLugo et al. 2004;Kee et al. 2016, 2018a,b). For S106 IR, the outflow cavity is also ionised, resulting in the striking bi-polar HIIregion (Schneider et al. 2018) and pointing to a more “evolved” stage compared to S140 IRS1, where only the disc is ionised (Hoare 2006).
Other high-mass sources, MWC349A (Báez-Rubio et al. 2014; Zhang et al. 2017) and Orion Source I (Matthews et al. 2010; Greenhill et al. 2013; Issaoun et al. 2017; Ginsburg et al. 2018) also drive disc winds. There is still some debate as to whether MWC349A is a pre- or post-main sequence star, while in Orion Source I the radio emission is not consistent with an equatorial ionised disc wind (Matthews et al. 2010). In these cases, the disc winds from these sources co-rotate with the disc (see alsoKlaassen et al. 2013). For MWC349A, the rotation is seen in maser emission (Zhang et al. 2017) as was also the first clear case of rotation in Orion Source I via SiO maser analysis (Goddi et al. 2009;Matthews et al. 2010) before it was recently found to be rotating in thermal emission as well (Hirota et al.
2017). Co-rotating disc winds are different from the aforemen-tioned equatorial wind scenario where the winds are radiatively driven, as they instead are thought to be intrinsically linked with the accretion process and disc (e.g.Turner et al. 2014). Identi-fying these disc winds in massive YSOs could therefore provide insight on accretion and mass loss processes.
In this article we report on our target G17.64+0.16 (here-after G17.64, also known as AFGL 2136, G017.6380+00.1566, CRL 2136, and IRAS 18196-1331). We originally targeted G17.64 along with five other luminous O-type (proto)stars1with ALMA and found mixed evidence for discs around the various targets using the typical CH3CN tracer (Cesaroni et al. 2017).
Our detailed investigations to-date of two of these six sources, G31.41+0.31 and G24.78+0.08 (Beltrán et al. 2018;Moscadelli et al. 2018), found clear velocity gradients in the CH3CN 12-11
line transitions. In G31.41+0.31, the gradient is indicative of rota-tional spin-up and infall towards the two central cores, whereas for G24.78+0.08, it is due to wind and outflow feedback from the O-star at the centre of an HCHIIregion. G17.64, presented in this work, is thought to be a massive O-type YSO at a later formation stage, but prior to forming a compact HIIregion.
G17.64 is located at 2.2 kpc and has a bolometric luminos-ity of 1 × 105L
. As part of the Red MSX survey2 (Lumsden
et al. 2013) G17.64 has the most up-to-date luminosity estimated using the detailed multi-wavelength spectral energy distribu-tion method from Mottram et al. (2011) updated to include fluxes from the Herschel HiGAL survey (Molinari et al. 2010), and combined with an unambiguous kinematic source distance (Urquhart et al. 2012,2014). The luminosity positions G17.64 as an O9−O9.5 type based onVacca et al.(1996) with an expected mass of 20 M , hence being selected as one of our six sources
to study O-star discs (Cesaroni et al. 2017). G17.64 is bright at near- and mid-IR wavelengths (2.2–24.5 µm −Kastner et al. 1992; Holbrook & Temi 1998; de Wit et al. 2009; Murakawa et al. 2013) as a reflection nebula that is co-spatial with the cavity walls carved by the blue-shifted lobe of the large-scale molecular outflow at a position angle of ∼135◦ (Kastner et al.
1994;Maud et al. 2015b) and perpendicular to a suggested polar-isation disc (Murakawa et al. 2008). Murakawa et al. (2013) analysed the IR spectrum from G17.64 and noted that the weak Brγ line is extemely broad (FW H M = 133 km s−1), indicative of an underlying wind. The surrounding molecular cloud has a mass of ∼600 M within a ∼1.2 × 1.2 pc region (Maud et al.
2015a). G17.64 also has a plethora of other observations rang-ing from the IR to radio (e.g.Minchin et al. 1991;van der Tak et al. 2000a,b;Menten & van der Tak 2004;Wang et al. 2007;
Lu et al. 2014) that image the bi-polar nebulae, the compact radio emission and surrounding environment in a range of typical molecular species such as NH3, CS, CH3OH at mm wavelengths.
Notably, G17.64 was also observed using mid-IR interferometry (de Wit et al. 2011;Boley et al. 2013). Models to reproduce the single MIDI baseline emission presented byde Wit et al.(2011) suggest G17.64 could have a compact disc. In the sample of mas-sive YSOs investigated byBoley et al.(2013) the combination of Keck data and multi-baseline MIDI observations also suggest a compact (<100 au diameter) dust disc is present in G17.64.
1 For consistency with Cesaroni et al. (2017) we use the same
nomenclature “(proto)star” only here. Throughout the paper we use young-stellar-object (YSO) to avoid any ambiguities as massive sources (>10M ) are probably buring deuterium and those at later stages
begin-ning to burn hydrogen (e.g.Hosokawa et al. 2010) and are not really true “proto”-stars.
2 http://rms.leeds.ac.uk/cgi-bin/public/RMS_DATABASE.
In Sect.2, we present the observation overview, while the main observational results are shown in Sect.3. In Sect.4, we analyse the data using a parametric model and also discuss G17.64 in reference to other known massive YSOs. Finally, the conclusions are given in Sect.5.
2. Observations
Our ALMA 12 m observations were conducted during Cycle 2 in July and September 2015 (2013.1.00489.S - PI: Cesaroni). The spectral setup included one wide-band spectral window (SPW) with ∼1.875 GHz bandwidth in order to obtain a continuum esti-mate and numerous other SPWs with higher spectral resolution and narrower bandwidths ∼234 MHz for the molecular line emis-sion in Band 6 (220−230 GHz). The individual datasets had between 38 and 41 antennas at any one time and baselines rang-ing from 40 out to ∼1500 m. We refer the reader toCesaroni et al.
(2017) for more details on the setup.
The data reduction was undertaken usingCASAversion 4.5.3 (McMullin et al. 2007). Owing to G17.64 having weaker molec-ular line emission than the other O-stars in the sample (cf.
Cesaroni et al. 2017) there were many line-free regions in all SPWs and thus a standard continuum subtraction was under-taken followed by self-calibration using the line-free continuum emission3. The self-calibration solutions were applied to the entire G17.64 data and the continuum subtraction was remade, thus both the continuum and line data have the self-calibration applied. The high signal to noise of G17.64 during each iterative stage of self-calibration allowed us to self-calibrate down to a minimal timescale of 6 s (i.e. starting with the source scan time of approximately 6 min, and subsequent iterations using 180, 60 and 30 s, then finally down to the integration time of 6 s). The improvement of the continuum signal-to-noise was almost a factor of five, improving from ∼180 to over ∼860.
A Briggs robust weighting of 0.5 (Briggs 1995) was used for all images to balance resolution and sensitivity, while a robust value of 1.5 was also used for the continuum in order to weight towards shorter baselines, sensitive to larger scale diffuse emission. Table 1 presents the beam sizes, noise levels and velocity resolution of the various images used in this work. We note that a subset of these data were first presented inCesaroni et al.(2017), namely, images of SiO and CH3CN, which were
analysed in a systematic fashion along with the other target sources. Here, however our data have undergone re-imaging and analysis after self-calibration. The self-calibration notably improved the signal-to-noise of our images and given the significant improvement in phase correction, have acted to sharpen the image structures and boost the detected fluxes (see also Cornwell & Fomalont 1999). Our results superseed those presented inCesaroni et al.(2017).
A stand-alone ACA and Total Power (TP) project, specif-ically aimed at understanding the outflows from our selected targets, observed G17.64 as part of Cycle 4 on 27 November 2016 (2016.1.00288.S - PI: Cesaroni). For the ACA data 11 anten-nas were available and the array had baselines ranging from 8.9 to 45.0 m. The total time on source was about three minutes for the ACA. Because of this disproportionately short on-source
3 TheSTATCONT package bySánchez-Monge et al.(2018) was used
for continuum subtraction in the other targets ofCesaroni et al.(2017) which have very few line-free channels. Comparisons were made for G17.64 usingSTATCONT which did produce comparable products although continuum subtraction in the image plane does not allow for self-calibration after the analysis.
time, the ACA data have a very sparse u,v coverage and little overlap with the 12 m baselines, and are of a low sensitivity. Therefore, the ACA data are unsuitable for merging with the aforementioned 12 m data in creating a representative image or well formed and weighted synthesised beam. The total power observations had ∼19 min on-source time and used data from three antennas in single dish mode. The TP pointing were set to cover the primary beam area of the ACA only ∼4500× 4500.
The spectral setup was selected to cover the same transi-tions as the previous 12 m array configuration and hence covered SiO and13CO, which are usually associated with outflow emis-sion. The data were pipeline reduced and then the ACA data were manually imaged using CASA version 4.7.2 (more recent data reduction uses a newer CASA version). The final beam size was ∼7.3800× 5.0400 with a position angle of 84.6◦ when
using a robust weighting of 0.5. The noise ranges from 0.2 to ∼1.0 Jy beam−1per 0.5 km s−1channel (the latter from channels
with significant emission). For the TP data we use the deliv-ered ALMA products. We note that the outflow parameters from these ACA data and a comparison to the TP data are presented in AppendixAand are not presented in the main paper.
3. Results
In this paper, we focus on the continuum and the kinematics of the SiO, while also presenting the emission from CH3CN,
CH3OH and H30α in order to describe the putative disc and
evi-dence for a disc wind. We also make a reference to the 13CO
observations from the 12 m array only, in order to describe the orientation of the molecular outflow “bubble”.
3.1. Continuum emission
Figure 1 shows the 1.3 mm dust continuum map of G17.64 imaged in our ALMA observations within a 300× 300region cen-tred on G17.64 (J2000 18h22m26.385s−13◦30011.9700). Due to
the high dynamic range, the contours at the 5 and 10σ levels best highlight the very weakest emission in the region. G17.64 is a very compact, bright source (peak = 74.2 mJy beam−1, flux density = 81.3 mJy) and is essentially isolated from other strong and compact cores. The strongest central emission com-ponent associated with G17.64 remains unresolved. The next strongest source in the field is a comparatively very weak (peak = 1.8 mJy beam−1) point source located 0.700(∼1540 au) to the south east, with another point-source further south and even weaker (peak = 0.7 mJy beam−1). As we achieve such a high sen-sitivity in the continuum map, we also detect diffuse continuum emission at the 5 to 10σ level and a number of point sources throughout the map, which are more clearly seen in Fig.2.
The central source, G17.64, is associated with compact radio emission (Menten & van der Tak 2004) that would contribute a maximum of 29.5 mJy (36% of the 1.3 mm flux density) when scaled to the 1.3 mm wavelength using the free–free spectral index of 1.2 estimated byMenten & van der Tak(2004) from fre-quencies <43.2 GHz. These authors also argue that the free–free spectral index will probably become shallower beyond 43.2 GHz, in which case if we assume a −0.1 optically thin free–free spec-tral index the contribution of the free–free at 1.3 mm would only be 3.57 mJy (4% of the 1.3 mm flux density). Thus, in the case where we assume optically thin millimetre emission and con-sider typical parameters for dust opacity coefficient, where κ0=
1.0 cm2g−1as suggested for densities of 106−108cm−3 for dust with thin ice mantles (Ossenkopf & Henning 1994), a gas-to-dust
Table 1. Noise, beam size and position angle, and velocity resolution for the 12 m array continuum and line data images made in this work.
Image Frequency Transition Eu Beam Noisea Resolution
(GHz) (K) (00,◦) (mJy beam−1) (km s−1) Continuum 226.700 ... ... 0.17 × 0.14, −88.2 0.07 2.65 (GHz) Continuum (robust 1.5) 226.700 ... ... 0.21 × 0.17, −87.9 0.16 2.65 (GHz) 13CO 220.39868 2−1 15.87 0.18 × 0.15,+89.8 <4.03b 0.35 SiO 217.10498 5−4 31.26 0.20 × 0.15, −88.3 1.21 2.70 CH3CN 220.70902 12−11 K= 3 133.16 0.18 × 0.15, −89.2 2.08 0.35 CH3OH 218.44005 4(2,2)−3(1,2)−E 45.46 0.20 × 0.15, −88.4 1.16 2.70 H30α 231.90093 Recomb. line ... 0.17 × 0.15, −90.0 1.58 1.40
Notes. A robust weighting of 0.5 was used unless indicated. For molecular lines the transition and upper energy levels are also listed. Units are as indicated in the column heading unless otherwise stated.(a)The noise is per channel for line data cubes.(b)The13CO has resolved out structure
leading to higher noise in channels where there is strong, poorly sampled emission. Line free channels have a noise of 1.76 mJy beam−1ch−1.
Fig. 1.ALMA band 6 (1.3 mm) continuum image of G17.64 made with
a robust 0.5 value. The contour levels are set at −3, 5, 10, 50 and 200 σ (1 σ = 0.07 mJy beam−1) to indicate the weak southern sources and the
strength of G17.64 (negative contours are dashed lines). The map is centred (0,0) on G17.64 at J2000 18h22m26.385s −13◦
300
11.9700
. The synthesised beam is indicated to the bottom left of the figure, while a scale bar is shown in the bottom right.
ratio of 100, and temperatures between 50 and 100 K, the mass of the compact dust emission ranges from 0.85 to 2.69 M when
considering the extremes of the free–free contribution (see also Eq. (1) ofMaud et al. 2013). The dust mass should be interpreted with caution as we resolve out emission with the interferometer, we do not correct for dust opacity, and we use a single common temperature. Considering these points together, uncertainties in dust mass by a factor of 2–3 are not unreasonable.
A two-dimensional Gaussian fit to the central continuum emission associated with G17.64 in the image plane results in deconvolved source dimensions (FWHM – full-width-half-maximum) of 53 × 39 milliarcsec (117 au × 86 au) at a position angle (PA) of 25 ± 5◦(east of north). When considering the ori-entation of the IR-reflection nebula and the large-scale 13CO
molecular outflow (see also Fig.2 and Appendix A) at a PA of ∼135◦, the deconvolved PA of the dust structure could
sug-gest that the mm emission is tracing an underlying disc in an almost perpendicular direction. The result is also consistent with
the IR-interferometric observations byBoley et al.(2013) where their Gaussian fitting suggested a disc FWHM of 44 milliarcsec-ond (97 au at 2.2 kpc) with a PA around 38◦. Although we have
very high signal-to-noise, we note that the strong centrally com-pact dust emission is essentially unresolved and therefore warn that the deconvolved parameters should be interpreted with some caution. Table2indicates the continuum parameters for G17.64. 3.2. Larger scales and the13CO molecular outflow
A large-scale view of the G17.64 region is shown in Fig.2. The background is a three colour composite image made using Ks (2.2 µm, red), H (1.6 µm, green) and J (1.2 µm, blue) images. These near-IR images were obtained by the UKIDSS (The UKIRT Infrared Deep Sky Survey;Lawrence et al. 2007) Galac-tic Plane Survey (Lucas et al. 2008) and have point spread function of ∼0.5−0.800, respectively (see alsoMurakawa et al.
2008). This is overlaid with grey contours at the 3, 4 and 5 σ lev-els indicating the 1.3 mm continuum ALMA image made with a robust 1.5 weighting. The outflow emission from13CO is shown by the coloured contours while the outflow direction is indicated by a white dotted line. G17.64 is identified as a plus symbol located at (0,0) in the map and, for reference, has a VLSR of
22.1 km s−1. The two sources to the south of G17.64 (as shown in Fig.1) are clearly seen, while there is also another compact source located in the region of blue-shifted13CO emission (at
a 2.8, −4.0 offset from G17.64), although only at the 4–5 σ level. There is diffuse continuum emission about 1.5 to 2.000
(3300−4400 au) to the north-east and south-west of G17.64, in a direction almost perpendicular to that of the outflow, suggestive of a remnant from a larger scale possibly dusty toroidal structure or filamentary dark lane (dashed white ellipse). Although our observations are sensitive to scales up to ∼700the diffuse
emis-sion has a low surface brightness (peak ∼5−10 σ) and appears disjointed, although we do not consider this to be real sub-structure. Far to the south-east (∼600, ∼13200 au) there is also another diffuse structure, again peaking at around the 10 σ level and of low surface brightness, such that we may only image the peak of even larger-scale diffuse emission. Lower resolution (∼0.800) shorter baseline observations better image the surround-ing extended, weak, dust emission and are consistent with our data in that we only detect the peaks of emission while larger structures are present (Avison et al., in prep.). The diffuse emis-sion we detect is located in the regions devoid of IR emisemis-sion. All other continuum features in our map, but not mentioned, are at the 3 σ level.
Table 2. Continuum observation parameters for G17.64+0.16.
Name Coordinates Wavelength Peak flux Flux density Free–free flux Massa
(J2000) (mm) (mJy beam−1) (mJy) (mJy) (M
)
G17.64+0.16 18h22m26.385s−13◦30011.9700 1.3 74.2 81.3 3.57−29.5 0.85−2.69
Notes. (a)Mass extremes estimated using dust temperatures between 50 and 100 K, a gas-to-dust ratio of 100, and a dust opacity coefficient of
1.0 cm2g−1.
In our observations the13CO emission serves as a tracer for
the outflow as we do not cover the more abundant typically used outflow tracer, 12CO, in our spectral setup. The 13CO clearly
highlights only the strongest emission regions in the outflow as the emission is over-resolved with our 12 m main array con-figuration. At the highest velocities, the blue-shifted emission (16.0–18.3 km s−1) peaks to the south-east while the red-shifted emission (35.0–41.0 km s−1) peaks to the north-west. This
high-velocity emission follows a roughly linear path through G17.64 and is highlighted by the cyan and orange contours, respectively, in Fig.2. The blue-shifted emission also appears co-spatial with the bright IR features to the south-east. The lower velocities (blue-shifted from 18.3 to 21.4 km s−1, red-shifted from 24.7 to 35.0 km s−1, recall the V
LSR is 22.1 km s−1), represented by
the blue and red contours trace the cavities resulting from a bubble-like shape created by the outflow. These appear to visu-ally match to the broadened outflow cavities shown inKuiper & Hosokawa (2018). The red-shifted emission appears to trace almost a tear-drop shape with an origin centred on G17.64, while the blue-shifted emission is only visible as an extended arc shaped structure far to the south-east of G17.64. The red-shifted bubble appears as roughly symmetric about the outflow axis and confined at the base within the diffuse dust emission either side of G17.64, meanwhile the blue-shifted arc is located inward (towards G17.64) of the diffuse dust emission to the south-east. The origin of the13CO emission appears to be from the working
surfaces of the outflow interacting with the surrounding diffuse dust in IR dark regions. All estimates of the outflow parame-ters are made using only the ACA data which are presented in AppendixA.
Notably, there is also overlapping blue- and red-shifted emis-sion located to the west of G17.64 (Fig.2, RA offset −0.500). In the context of an outflow this can only be explained if there is a wide cavity opening angle such that the flow close to G17.64 is emitted almost radially in the plane of a disc, and thus has both blue- and red-shifted components close to the star before the flow is redirected to follow the parabolic shape of the cavity where the bulk motion becomes overall red-shifted to the north-west (e.g.
Arce et al. 2013). 3.3. SiO
Figure 3 (left) shows the moment zero map of the SiO (5–4, Eu= 31.26 K) emission integrated over the central ∼3 km s−1of
the line covering the VLSR, between 21.4 and 24.0 km s−1. In
contrast to Cesaroni et al. (2017), here we integrate the SiO emission over a narrower velocity range to highlight the elon-gated structure at near VLSRvelocities. The elongated structure
of the emission is clearly visible and perpendicular to the large-scale outflow (Fig.2). Fitting the integrated SiO emission with a 2D Gaussian in the image plane indicates a PA of ∼47 ± 7◦ and a major axis of ∼0.2800 (∼600 au). The deconvolved PA is
30 ± 10◦consistent with the PA of the dust continuum emission.
The SiO emission is broad in velocity, ranging from −3.0 to 45.7 km s−1considering emission above 3 σ, this is broader than
presented inCesaroni et al.(2017), 4–39.1 km s−1, owing to our improved images after self calibration. In most massive YSOs, SiO is very spatially extended, typically tracing emission at shock fronts (e.g.Schilke et al. 1997) created by active jets and outflows driven by the energetic central sources (e.g.Cabrit et al. 2007;Sánchez-Monge et al. 2013b; Duarte-Cabral et al. 2014;
Klaassen et al. 2015; Cunningham et al. 2016; Cesaroni et al. 2017;Beltrán et al. 2018;Moscadelli et al. 2018). In at least one massive YSO, Orion source I, the SiO emission is rather com-pact, associated with SiO masers and is shown to trace a rotating disc and disc wind structure (Goddi et al. 2009;Ginsburg et al. 2018). The SiO emission from G17.64 is also relatively com-pact, unlike a jet. There is also compact emission from other silicon and sulphur bearing species in our observations, emis-sion from SiS (v = 0 12–11, Eu = 69.95 K) and33SO (65–54,
Eu = 34.67 K) are centrally peaked at the location of the
con-tinuum and SiO emission. The strong detection of SiO points to a reservoir of silicon in the gas phase, which, in addition to the detected sulphur-bearing species, suggest an association with ionised, hot (>100 K), turbulent and shocked gas (e.g.Minh et al. 2010;Minh 2016) where these species are released from the grain material.
The solid contours in Fig.3(left) represent the moment zero emission integrated over the highest blue- and red-shifted veloc-ities of the SiO emission (between −3.1 and 10.6 km s−1 and
32.0 and 45.7 km s−1, respectively) whereas the dashed lines
out-line the emission integrated over only the low velocity emission (between 14.4 and 21.4 km s−1and 24.0 and 29.7 km s−1 for the blue- and red-shifted emission, respectively). The PA of a simple line drawn through the centroids of the corresponding blue- and red-shifted emission at the high velocity extreme is consistent with that from the integrated map over the central velocities and with the PA at ∼30◦ of the continuum emission. Observation-ally, the SiO emission from G17.64 appears inconsistent with a jet or outflow origin when considering the known CO outflow and the IR nebulosity at a PA of ∼135◦ and the lack of any
evidence for a jet with an axis close to the putative disc PA of ∼30◦ (Kastner et al. 1992;Murakawa et al. 2008). Rather, the
SiO appears to originate from a disc-like structure perpendicular to the outflow. Intriguingly, the spatially extended lower velocity blue- and red-shifted emission (dashed contours) appear offset in the opposite direction from the spatially compact high veloc-ity blue- and red-shifted emission (solid contours). If we were to assume a rotating disc is present, the high (±20−25 km s−1) and low (±5−10 km s−1) velocity rotational directions disagree,
that is, with one suggesting a small disc rotating one way – with blue-shifted emission to the north-east and the other indicat-ing a larger disc in the opposite direction – with blue-shifted emission to the south-west (see below and Sect.4.1). In the right panel of Fig.3, we show the centroid positions after fit-ting Gaussians to each channel in the SiO image cube between
Fig. 2.Three colour near-IR composite image made using Ks (2.2 µm, red), H (1.6 µm, green) and J (1.2 µm, blue) images obtained by the UKIDSS overlaid with grey contours of the ALMA band 6 (1.3 mm) continuum emission (robust 1.5 weighting) from the 12 m array only. The contours are at the 3, 4, 5 σ levels (1 σ = 0.16 mJy beam−1) to highlight the weaker emission features. The position of peak continuum emission from G17.64
is indicated by the plus symbol at (0,0) offset (J2000 18h22m26.385s−13◦
300
11.9700
) – the continuum contours are not plotted due to the overlap with13CO contours. In the IR, G17.64 at the centre is the dominant source, while to the east, west and south-east the emission is associated with
the reflection nebula. The white dashed ellipse is indicative of a possible remnant dusty toroid, or underlying dark lane structure. The coloured contours show the moment zero maps of the13CO outflow emission integrated over various velocity ranges. The cyan and orange contours show
the blue- and red-shifted highest velocity emission integrated from 16.0 to 18.3 km s−1and 35.0 to 41.0 km s−1, respectively. The contour levels are
at 5, 10 and 15 σ, where σ is 3.28 and 4.57 mJy beam−1km s−1for the blue- and red-shifted high velocities. The blue and red contours indicate the
integration of the lower outflow velocities from 18.3 to 21.4 km s−1for the blue-shifted emission and from 24.7 to 35.0 km s−1for the red-shifted
emission. We note that the source VLSR= 22.1 km s−1. These contours are at 4 and 5 σ, where σ is 6.25 and 10.06 mJy beam−1km s−1for the
blue-and red-shifted velocities. The dotted white line indicates the outflow axis at a PA of ∼135◦
(Kastner et al. 1994;Maud et al. 2015b, see also AppendixA) which passes though G17.64 and the points of highest velocity blue- and red-shifted outflow emission along with the IR nebula. A synthesised beam for the mm emission is shown to the bottom-left, a scale bar to the bottom-right, and the legend for the integrated velocities at the top. The primary beam cut off is beyond the map extent.
−0.3 and 45.6 km s−1. It is clear that the strongest emission com-ponents of each channel are within 0.0300of the central location
of G17.64, and moreover that the blue- and red-shifted emission are offset to the north-east and south-west, respectively, support-ing the idea of a rotatsupport-ing structure. The typical errors on the fits are of the order half a pixel, corresponding to 0.007500. As the
strongest emission components from all channels match the con-ventional rotation of a disc, the extended emission preferentially seen at low velocities and appearing to show opposite rotation due to the spatial extent, probably originates from a kinemati-cally separate structure that we cannot fully disentangle in the observations.
Fig. 3.Left panel: moment zero image of the SiO emission integrated over the VLSRbetween 21.4 and 24.0 km s−1. The grey contours highlight the
elongated emission and are at the 10, 15 and 20 σ levels (σ = 5.524 mJy beam−1km s−1). The blue and red solid contours show blue- and red-shifted
high velocity SiO emission thought to be tracing the high velocity component, possibly an underlying rotating disc (levels at 10, 15 and 20 σ, where σ = 9.57 and 9.41 mJy beam−1km s−1for the blue- and red-shifted high velocity emission) while the dashed lines indicate the 10 σ level (σ = 7.54
and 6.58 mJy beam−1km s−1for the blue- and red-shifted emission) attributed to the lower velocity component. The velocity ranges are indicated to
the top of the figure, a scale bar to the bottom right, and the synthesised beam to the bottom left. The black solid line is the direction of the outflow at a PA of 135◦
and the dotted is that of the disc major axis at a PA of 30◦
, corresponding to the cut used for the PV analysis (see Fig.4). Right panel: plot of the centroid positions from a 2D Gaussian fitting of each channel from the SiO image cube between −0.3 and 45.6 km s−1, clearly
showing that the blue- and red-shifted emission are predominantly to the north-east and south-west, respectively. We note that the scale is changed (by a factor of ten) compared to the left panel. A scale bar is also indicated to the bottom right.
In Fig.4, we show the position-velocity (PV) information of the SiO emission extracted from the image cube along a cut at a PA of 30◦, that is, along the major axis direction of the
puta-tive disc from the north-east to south-west direction. We include all data along the cut within a width equal to that of the synthe-sised beam, which helps to slightly improve the signal-to-noise. Overlaid on all images are dashed grey lines to cut the four quadrants. The left image includes grey contours to indicate the general structure of the emission along with a black dotted line to guide the eye as to the skew of the PV diagram indicating high velocities at small offsets (±0.100). Such a shape in the PV plane can generally be interpreted as an unresolved rotating structure, although we additionally see considerable large-offset emission at low velocities (“spurs”) at both blue- and red-shifted veloci-ties (see Sect.4.1). We notice that there appears to be a “kink” in the lowest level grey contour that could be interpreted as a sep-aration between different components, such as rotation (close to linear structure highlighted by the black dotted line) and radial motion, infall or expansion (represented by a diamond like shape, e.g.Tobin et al. 2012).
For illustration we over-plot lines of Keplerian rotation for various central masses (cf.Cesaroni et al. 2014) to explain either the high velocity, small offset emission from a compact disc around a 20, 30 or 40 M mass O-type YSO, or the lower
veloc-ity, large offset emission as a larger disc rotating in the opposite sense about a 10, 20, or 30 M mass OB-type YSO (central and
right panels of Fig.4, respectively). We used a source inclina-tion of 70◦, where 90◦is a view edge-on to the disc, based upon
models byde Wit et al.(2011). We do not suggest that there are two counter-rotating discs in G17.64, but present both separate scenarios. Furthermore, we note that our data only marginally resolve the structures in the PV image compared to studies of low-mass YSOs and as predicted for ALMA high-resolution (sub-100 au) observations from models (e.g. Yen et al. 2014;
Harsono et al. 2015;Seifried et al. 2016;Dutrey et al. 2017) and thus “fitting” Keplerian rotation lines is not reliable. It is clear that any of the Keplerian rotation curves only represents emis-sion in two of the four quadrants in the PV plots and thus a disc alone cannot fully describe the data, independent of disc size or rotation direction. Furthermore, pure radial motion, infall or expansion, would result in a symmetric diamond shaped PV plot (seeOhashi et al. 1997; Brinch et al. 2008; Tobin et al. 2012;
Sanna et al. 2018) which alone would not represent the skew seen to high velocities at small offsets. A combination of kinematics components is required. A more detailed analysis, taking also the observational limitations (spatial and velocity resolution) into account is presented in Sect.4.1.
3.4. CH3CN and CH3OH
CH3CN is one of the most commonly used tracers to study the
kinematics of small scale (<0.1 pc) dense gas surrounding mas-sive YSOs. It is specifically used to investigate the presence of disc signatures and can also probe kinetic temperatures using the various emission lines from the close in frequency “K” rota-tional ladders (Araya et al. 2005; Purcell et al. 2006; Beltrán et al. 2014; Cesaroni et al. 2014; Sánchez-Monge et al. 2014). For G17.64 we detect the J = 12–11, K-ladder from K = 0 to K= 7, Eu= 68.86−418.63 K (although only at the 2–3 σ level for
the highest K-lines, K = 5−7). In correspondence withCesaroni et al.(2017, Fig6), who plot only a relatively cooler K = 2 line, we find no evidence for a rotating disc in this molecular tracer, as the CH3CN emission is preferentially detected only to the
south-west of the continuum peak of G17.64 (within 100, 2200 au)
throughout all velocities.
Figure 5 shows the channel map from the CH3CN (J =
12–11) K = 3 line (Eu = 133.16 K) within a 3.0 × 3.000 region
Fig. 4.Left panel: position-velocity (PV) diagram of the SiO emission extracted from a cut (as wide as the synthesised beam) at a position angle of 30◦
across the putative disc major axis (see Fig.3– position offsets are positive to north-east). The grey contours are at the 10, 15, 20, 25 and 30 mJy beam−1level to highlight the structure. The black dotted line is a guide to indicate the excess of high velocity emission at small spatial
offsets. Centre panel: observed SiO PV as the left plot (colour map with grey contour at a 10 mJy beam−1level) overlaid with pure Keplerian
rotation profiles for a 200 au radius disc around a 20, 30 and 40 M (inner dashed, solid, outer dashed lines) O-type YSO, when considering an
inclination of 70◦
(de Wit et al. 2011). The disc outer radius was limited to 200 au such that the lowest velocities roughly matched the kink in the observed data. The disc lines here represent only the high-velocity emission. Right panel: observed SiO PV as the left and central plots but overlaid with the contours of pure Keplerian rotation profiles for a disc around a 10, 20, 30 M (inner dashed, solid, outer dashed lines) OB-type
YSO rotating in the opposite direction to that plotted in the central panel (using an inclination of 70◦
). The disc outer radius was set to 650 au in an attempt to match the low velocity spurs at large spatial offsets. The spatial and velocity resolution are indicated by the black bars at the bottom left of all plots. Note the central figure indicates the north-east and south-west offset directions that relate to the cut along the disc direction in Fig.3.
A plume shaped structure is evident at marginally blue-shifted velocities (19.7−22.5 km s−1). At these velocities all low-K tran-sitions (K = 0−4, Eu = 68.86−183.15 K) show the curved
“plume”-like shape. A “knot” feature at the tip of the plume coincides with the inner edge of the diffuse continuum dust emission 1–200 (2200−4400 au) south-west of G17.64, possibly an emission enhancement due to interaction with the surround-ing material, the dusty toroid or dark lane structure (see the continuum contours in the 20.7 km s−1 panel, Fig.5). Emission
from the plume is co-spatial with the13CO (see Fig.2),
poten-tially from a wide-angle flow driven by G17.64, and supports such an interpretation. Indeed, if a flow was following a cavity with a wide opening angle, we would also expect some spatial overlap of the blue- and red-shifted emission close to the source. Furthermore, the blue-shifted emission from K < 3 lines has a distinctive arc of emission to the south-east that is at the leading edge of the blue-shifted arc already identified in13CO, from the interaction of the larger scale outflow bubble with the surround-ing medium (this is not shown in Fig.5but the CH3CN matches
the13CO arc structure in Fig.2to the south-east).
In Fig.6 (left) we present the moment zero map of the CH3OH line at 218.440 GHz (4(2,2)−3(1,2)–E, Eu= 45.46 K)
inte-grated between 19.0 and 23.9 km s−1 to highlight the plume structures curving away from the plane of the assumed dust continuum and SiO disc (PA ∼ 30◦ – faint dotted line). Both plumes are coincident with the over-plotted OH masers detected byArgon et al.(2000), not only spatially, but also in velocity. The north-east OH maser is at 23.3 km s−1while those to the south-west are between 19.9 and 20.9 km s−1. Multiple other CH
3OH
(including13CH3OH) lines are detected in the region
surround-ing G17.64, but all are weaker. In the right panels of Fig.6, we also show the moment zero maps of CH3CN (J = 12–11) K = 3
(Eu = 133.16 K) and K = 4 (Eu = 183.15 K) integrated over the
same range as CH3OH (between 19.0 and 23.9 km s−1), which
also indicate the plume structure. The K = 4 emission is notably weaker than the K = 3 line of CH3CN.
The west plume like structure is mostly associated with blue-shifted emission as is the aforementioned arc-like structure to the
south-east, seen in emission from CH3OH, 13CO and CH3CN.
Although weaker, a second plume like shape to the north-east of the continuum peak appears to be positioned as an almost inverse-mirror-image of the westerly plume, but it is only evi-dent at slightly red-shifted velocities in the CH3OH transition at
218.440 GHz. Comparing to the CH3CN K = 3 emission, we find
that there is a faint emission spot coincident with the strongest region of CH3OH in the eastern plume (Fig.5, 23.2–23.9 km s−1,
offset +1.500, +0.400) − close to the OH maser spot.
The combination of the emission and velocity structure of these tracers and the OH maser emission all point to them tracing the cavity working surfaces where the wide-angle wind interacts with the inner edge of the dusty structure around G17.64. 3.5. SiO and CH3CN connection
In Fig.7, we show the PV diagram of the CH3CN (J = 12–11)
K= 3 emission extracted along the same cut as that for the SiO (NE to SW, see the disc major axis in Fig.6), however we use a width of 12 beams in order to encompass emission from the plume and the knot. The majority of the emission is at negative offsets, to the south-west, with only slight blue-shifted emission to the north-east (+0.2 offset). From an offset of −0.1 to −0.300
we see co-spatial blue- and red-shifted emission, as we might expect if there were a wide opening angle cavity where mate-rial flows almost along the line-of-sight closest to the source. Conversely, it could also be indicative of simultaneous infall at both the front and rear of the disc. However, infall would gen-erally be understood to speed up as it comes closer to the star, whereas the brightest and strongest emission features from the plume indicate the opposite, the velocities (both blue- and red-shifted) increase from an offset of about −0.0800to an offset of −0.2800. Thus, the wind scenario, where material is continually driven away from the star in the plane of the disc, could gen-erate such a PV diagram. Note, the PV diagram does indicate an overall triangle shape considering all the emission and shows generally higher velocities closer to G17.64, and could be inter-preted as an infalling stream. However, this assumes that all of
Fig. 5.Channel map of the CH3CN (J = 12–11) K = 3 line ranging from 17.9 to 26.3 km s−1in steps of 0.3 km s−1. The contours of the CH3CN
emission are at the 3, 5 and 10 σ level (σ = 2.0 mJy beam−1chan−1). In the central panel at the V
LSRvelocity channel (22.1 km s−1) the dotted line
indicates the PA of the disc major axis at 30◦
, while the solid (blue and red) line indicates the outflow direction at a PA of 135◦
. At all velocities the CH3CN emission is brightest to the south west of the main continuum peak (at 0,0 indicated by the “+” symbol). The plume structure (indicated)
is visible between ∼19.7 and ∼22.5 km s−1while the knot at the tip of the plume is inward of the diffuse dust emission, potentially indicative of an
interaction of outflowing material with the surrouding dust structures. The continuum emission at the 3, 4, 5 σ levels as in Fig.2is plotted in black contours in the 20.7 km s−1channel panel to highlight the diffuse dust emission. A scale bar and synthesised beam are shown in the bottom left
panel.
the emission is from a spatially coherent structure, whereas in the channel map for CH3CN (Fig.5) there is a physical break
between the plume and knot. The knot emission spatially offset at −0.500is seen at bluer velocities and is disjoined from the main emission. Rather than being part of an infalling flow or stream, that contradicts with the plume only emission which increases in velocity with distance from G17.64, the knot is potentially emission from another structure, or given its location, due to an over-density where the wind impacts the proposed dust struc-ture around G17.64. All the CH3CN emission is within the low
velocity spur regions at larger positional offsets when compared with the SiO emission. The CH3CN could be indicative of the
interface between a wide-angle wind blown from the putative disc with the surrounding diffuse medium.
3.6. H30α
H30α hydrogen recombination line emission (231.9 GHz) is detected from G17.64. It appears weaker than that from G29.96, G24.78 and G345.49, the other O-stars in our sample where it is also detected (Cesaroni et al. 2017). The emission from G17.64 is exceptionally broad (σ ∼33.5 km s−1, or FWHM ∼ 81.9 ± 1.7 km s−1from a Gaussian fit – Fig.8) and spatially
Fig. 6.Left panel: moment zero map of the CH3OH line at 218.440 GHz integrated between 19.0 and 23.9 km s−1(colour map). The contours of the
CH3OH emission (grey) are at the 3, 5, 10, 15 and 20 σ level (σ = 7.44 mJy beam−1km s−1), while those of the continuum emission at the 3, 4 and
5 σ levels are also shown in black (see also Fig.2). The plume emission both to the east and west curves away from the plane of the proposed SiO disc plane (black dotted line) possibly due to a redirection or collimating of a wide-angle flow to the common CO outflow direction (blue and red solid line). The plus symbols mark the OH masers fromArgon et al.(2000). G17.64 is centred at (0,0). Right panel: moment zero maps of CH3CN
(J = 12–11) K = 3 (top panel) and K = 4 (bottom panel) also integrated between 19.0 and 23.9 km s−1 (colour map). The contours are at the
3,5,10,15,20 σ levels (σ = 4.93 mJy beam−1km s−1for the K = 3 transition and 7.61 mJy beam−1km s−1for the K = 4 transition. The dotted line in
all panels is that of the SiO disc plane.
level. G17.64 can be regarded as a broad-radio-recombination-line source (c.f.Sewiło et al. 2008;Kim et al. 2017). We note that by eye there appears to be a narrow peak of emission at around 50 km s−1which could tentatively be identified as HNCO v = 0 28(1,2)–29(0,29), Eu∼ 470 K. Given the narrow width we do not
consider it to significantly effect the broad line-width found for the H30α line.
Following the analysis outlined in Galván-Madrid et al.
(2012, with electron and ion densities of 107cm−3 and a tem-perature of 9000 K) pressure broadening can at most explain ∼0.6 km s−1, while thermal broadening contributes ∼20 km s−1. Thus, if the line width we measure has only a typical small non-thermal turbulent contribution (∼5 km s−1, Sewiło et al. 2008), there must be a significant underlying bulk gas motion component (e.g. Keto et al. 2008), potentially comprised of infall (accretion), outflowing material (wind or expansion, cf.
Moscadelli et al. 2018) and/or rotation. We cannot rule out a combination of all mechanisms contributing to the bulk motion when we consider the emission is spatially unresolved in the cur-rent data. The high velocities and signpost of ionisation and high temperatures ties with the release of silicon and sulphur from the grains close to G17.64 due to associated shocks. The H30α emission will be addressed further in Klaassen et al. (in prep.).
4. Analysis and discussion
4.1. Modelling the kinematics of the compact disc emission We created a simple parametric model to further investigate the kinematics of the potential disc seen in SiO emission. This fol-lows in the spirit of the models outlined in Richer & Padman
(1991) where a synthetic PV diagram is produced for comparison with the observations (see alsoWang et al. 2012;Girart et al. 2017). From the outset, given the limitations of the data in terms of spatial resolution and due to degeneracies with various param-eters we do not aim to specify and fit every possible parameter, but rather address the qualitative physical structures and kine-matics that are consistent with the data and their interpretation. We make a qualitative by-eye assessment for suitable matches with the data. The main parameter we aim to characterise is the stellar mass, as this is key to understanding whether G17.64 is truly an O-type YSO.
We build our model in a 3D environment where each volume element holds an emissivity (ε(x, y, z)) and a velocity (V(x, y, z)). We only assume an optically thin representation, and consider only two geometric structures in the model, a thin disc and a thin disc surface layer. The emissivity of the structures are parameterised by power-laws representative of the combination of density and temperature with radius, ε ∝ (r/r0)p+q, where r
is the cylindrical radius, r0is an arbitrary normalisation radius,
and p+ q is the “combined” power-law index (e.g. Brinch & Hogerheijde 2010; Fedele et al. 2016), where p is the density power-law index and q is the temperature power-law index. The disc, or disc surface can be independently parameterised by its inner and outer radii. The relative strength, or weighting, of the emission from each structure is also a free parameter, εwei. The
models are limited to two velocity field descriptions: (1) Keple-rian rotation, Vrot = (GM/rd)1/2, where Vrotis the velocity as a
function of radius (rd) in the disc plane, G is the gravitational
constant and M is the stellar mass; and (2) radial motion in the disc plane. The magnitude of the radial motion is parameterised by the weighted free-fall velocity, Vrad = Awei(2 GM/rd)1/2,
Fig. 7.Position-velocity (PV) diagram of the CH3CN (J = 12–11) K = 3
line (colour map) overlaid with the PV diagram of the SiO as in Fig.4 (black contours). There is notable emission to negative spatial offsets related to a south-west direction along the PV cut, whereas to posi-tives offsets (north-east) there is barely any emission. Both blue- and red-shifted emission coincide to the south-west, related to the signature of either a wide-angle wind, with flows nearly radially, or that there is simultaneous infall from material in front and behind the edge of the disc. The emission at bluer velocities and high negative offset (−0.500
) is associated with the knot structure identified in Fig.5. The spatial and velocity resolution of the CH3CN map is shown at the bottom left, while
the direction along the PV cut to the NE and SW is indicated at the top.
Fig. 8.Spectrum of H30α line extracted from within a single beam
cen-tred upon the continuum peak location over-plotted with a Gaussian fit (dashed black). The velocity resolution is 1.4 km s−1. The line is centred
at 22.1 km s−1the V
LSRof G17.64 as indicated by the vertical grey line.
The base flux level for the Gaussian fit was fixed to zero. The line is weak compared to the other O-star sources (Cesaroni et al. 2017) but is exceptionally broad, with a FWHM of ∼81.9 ± 1.7 km s−1.
Table 3. Parameters for the various models tested, the parameter ranges and the increments used.
Parameter Range Increment
Disc - rotation model
Stellar mass M 5−50 10a
Disc inner radius Rinner 10−50 10
Disc outer radius Router 100−600 100
Emissivity powerlaw p+ q −2.6 to −1.8 0.2
System inclination i 30−90◦ 20◦
Disc - rotation and radial model
Stellar mass M 10−30 10
Disc inner radius Rinner 20−50 10
Disc outer radius Router 100−600 100
Emissivity powerlaw p+ q −2.6 to −1.8 0.2
Radial vel. weighting Awei 0.5−1.0 0.25
System inclination i 70◦ Fixed
Disc - rotation and disc surface radial model
Stellar mass (M ) M 10−30 10
Disc inner radius (au) Rinner 20−50 10
Disc outer radius (au) Router 100−600 100
Disc surf. inner radius (au) Rsinner 50−600 100b
Disc surf. outer radius (au) Rsouter 650 Fixed
Emissivity powerlaw p+ q −2.6 to −1.8 0.2
Emissivity weighting εwei 0.5−2.0 0.5
Radial vel. weighting Awei 0.5−1.0 0.25
System inclination (◦) i 70◦ Fixed
Notes. (a)Between 5 and 10 M
the increment is 5 M and 10 M
thereafter.(b)Between 50 and 100 au the increment is 50 au, and 100 au
thereafter.
where Aweivaries from 0.5 to 1, where 1.0 corresponds to pure
free-fall.
Our models are built from different combinations of struc-tures and kinematics, as detailed in the sub-sections below. The principle is that we begin with the simplest possible descrip-tion before including more complex scenarios with more free parameters.
For each volume element we establish the intensity of the emission by weighting with the emissivity profile, while the velocities are broadened by a Gaussian filter with a typical tur-bulent line width of 0.5 km s−1at disc scales (Richer & Padman
1991), although later we account for our observation spectral resolution. The model is inclined as required (an edge-on disc would be at 90◦) and the line of sight axis is integrated to make a synthetic datacube (RA, Dec, VEL). The cube is finally con-volved with a Gaussian kernel accounting for the spatial and velocity resolution before we extract the PV diagram. For com-parisons with the observations, we scale the peak emission of the model PV diagrams to that of the data. Table3indicates the var-ious model parameters and the possible ranges investigated. We explore a large parameter space with a relatively coarse model grid, again due to the limited resolution of the data at hand. Figure9 shows the representative “best” matches to the kine-matics of the observations from each model regime. The left plots indicate the PV plots as a schematic while the right plots show the models compared with the observations, top to bottom follows the models as presented in the following sub-sections.
Fig. 9. Various models tested as a schematic (left panel) and then as an overlay on the observed SiO data (right panel). From top to bottom: disc model with only rotational velocities; the disc model with rotation and radial veloc-ity motion; and the disc with disc sur-face model where the disc has rota-tional motion whereas the surface has radial motion. Both disc only models (top and middle panel) have intensity profiles where most emission is shifted to high velocities and small spatial off-sets, unlike the centrally peaked obser-vations. For the disc model with rota-tion and radial morota-tion (middle panel), the observer sees the vector sum of the velocity component, such that there is always a resulting velocity offset at larger spatial offsets due to the rota-tional component. This does not match the data. For models of a rotating disc combined with a disc surface that has radial motion (bottom panel), the veloc-ities from each structure are super-posed and thus can represent the entire observed PV diagram (the bottom-left plot shows the Keplerian and radial velocities profiles in the dotted thin line). In all right hand plots the data resolution is shown to the bottom-left while the contours for the data and model are from 20 to 80% in 10% steps of the peak emission. Note the basic models parameters shown in the right plots.
4.1.1. Disc only – rotation
Using models where rotation is the only velocity component we are able to match only two opposite quadrants in PV space, exactly as in Sect.3.3, even when accounting for the resolution of the observations (Fig.9, top panel). We can match the high-est velocities in the PV diagram using a small radius disc, but this model does not account for the low velocity spurs. A large outer radius disc with a respectively larger inner radius (vs. a small disc) rotating in the opposite sense does match the spurs, but cannot describe the high velocity emission. A considerable number of disc-only models represent the data (in two of the four quadrants) equally well (or equally poorly) primarily due to the degeneracy between stellar mass and inclination angle. Given the modelling undertaken previously byde Wit et al.(2011) we choose to fix the inclination angle to 70◦(where 90◦is edge-on
to the disc). A side effect is that the stellar mass is a lower limit
in all models considering the close to edge-on orientation of the disc.
With a fixed inclination the aforementioned high velocities are represented by discs with Rinner = 20–50 au and Router =
100–600 au for stellar masses ranging from 10 to 30 M in
var-ious parameter combinations. For larger masses the velocities at all radii become too great, and although we could increase the inner radius beyond 50 au the distribution of emission shifts away from low spatial offsets which is inconsistent with the data. The lowest mass of 5 M cannot provide sufficiently high
veloc-ities even with a 10 au inner radius (and such a low mass source is anyway inconsistent with the luminosity of G17.64). We there-fore rule out M < 5 M and M > 30 M stellar masses. There
is however still a degeneracy between mass and Rinner.
Con-volution of the model with the observable spatial and velocity resolution acts to reduce the high velocity emission peaks by smoothing the emission and also results in higher mass larger
inner radius models appearing the same as lower mass lower inner radius models. A disc rotating in the opposite direction, matching the spurs, is well represented by a more limited range where M = 10 M , Rinner = 20 au, and Router = 400−600 au. A
larger stellar mass in this case results in too high velocities at large spatial offsets. Variation of the emissivity power law has a marginal effect on the distribution of the intensity in the PV diagram. Steeper profiles (−2.6 to −2.2) produce more emission from smaller radii that have the largest velocities and thus begins to create two emission peaks, inconsistent with the observed intensity distribution. None of these models match the entire PV plot.
4.1.2. Disc only – rotation with radial motion
In an attempt to reproduce the emission in all four quadrants, we now include radial motion, interpreted as accretion through, or expansion of the disc, along with rotation. Each volume element in the model now has an assigned radial and rotational veloc-ity component (Fig.9, middle panels). Following from the above models we limit the mass range between 10 and 30 M . We find
a wide range of parameters that somewhat match the observed PV diagram partially, however, the disc is required to have an outer radius (Router) of at least 600 au in order to fit the spurs.
The addition of the radial motion actually narrows the model parameter space, requiring an inner radius between 30 and 50 au to better reproduce the observables, although it is still degenerate with stellar mass. The weighting of the radial to rotating veloc-ity components provides an additional degeneracy, although the parameter range of the “best” matching models is already broad. With a fixed outer radius Router of 600 au, the 10 M stellar
mass models provide a reasonable match for Rinner = 30–50 au,
and various combinations of disc power-law profile and radial to rotational velocity weightings. For a 20 M stellar mass, the
parameters are further restricted, Rinner = 40–50 au, power-law
profile > −2.2, and Awei= 0.5−0.75, while for a 30 M central
mass there is only one parameter set that matches the observed PV, Rinner= 50 au, power-law = −1.8 and Awei= 0.5.
Although these models have a more restrictive range, owing to the inclusion of radial motion compared to the rotation only models, the major problems are the considerable rotational velocities up to 5−6 km s−1 at the largest spatial offsets in the
spurs (±0.2300, i.e. >500 au radii) and the intensity distribution at low spatial offsets being biased to larger velocities than indicated in the observations. What we do deduce is that the high velocity components matching the skew in the PV diagram can only be provided by a disc rotating in one direction (e.g. Fig.4, left and centre), as a large disc rotating in the opposite sense (previously matching only the spur emission – e.g. Fig.4, right) is ruled out as the inner radius is too large in these cases and always under-represents the high velocities while still not fully matching the spurs. Again, none of the models are entirely consistent with the observed PV diagram.
4.1.3. Disc and disc surface
We now consider that the large (>400 au) spatial offset emission close to VLSRvelocities must be provided by a structure separated
from the rotating disc. We model this by including a thin disc surface with only radial motion with a variable weighting, as outlined above (Fig.9, bottom). This motion can be interpreted as either accretion onto the disc or a wind blowing away the disc. More realistically, there is likely a stratification of the disc in scale height where the velocity transitions from rotation to radial
motion, although constraining this detail is beyond the scope of our modelling. Here we fix the scale height of the surface layer to only one resolution element above and below the disc (10 au in our models) and note that is spatially separate from the thin-disc (in the mid-plane). This thin surface layer is the limiting case as we cannot constrain scale height information with the data at hand. Because the structures are separate, the corresponding kinematics are superposed in the final model PV diagram (Fig.9, bottom-left). The emissivity profile of the disc surface is exactly the same as the disc, although it can be weighted by a factor between 0.5 and 2.0, εwei.
There are obvious degeneracies in these models given the additional flexibilities of the independent disc inner and outer radii and disc surface inner radius along with the weightings of intensity and velocity components. We fix the disc surface outer radius to 650 au to represent the large offset emission near VLSR
velocities. Again, a range of models match the observations. For a 10 M stellar mass the models are limited to only Rinner= 20 au
as they otherwise do not reproduce the high-velocity emission, while Router < 200 au so as not to spread too much emission
to the lower velocity regions at larger radii (i.e. to the top-right and bottom-left quadrants). The disc surface inner radius, Rsinner, has to be <200 au in general so as to provide emission
from radial motion in the top-left and bottom-right quadrants of the PV plot. Increasing the stellar mass to 20 M we find the
same degeneracies as before, the inner radius of these models is Rinner = 30−40 au while Router < 200 au and Rsinner < 200 au,
as for the 10 M case. At the final mass tested, 30 M , again
Rinner must increase to 40−50 au, Router must remain <200 au
while Rsinneris between 100 and 400 au in general, but not well
constrained. Notably, for this highest stellar mass the match with the data appears slightly worse, primarily as the slope between high-velocity small-offset emission is slightly steeper compared with the lower mass models. For all the models, in general the other parameters can span essentially the full ranges, although importantly all models provide a qualitative match to the entire observed PV plot in all quadrents, depending on the parameter combination.
Overall the 10–20 M stellar mass models best represent the
data, although a number of 30 M stellar mass models are still
reasonable. We do rule out masses between 40 and 50 M as
they have too large velocities at all spatial offsets, as discussed in the disc only models, although we cannot further constrain the stellar mass. Comparing with the disc only models, the disc and disc surface models do provide a qualitative match to the entire PV diagram, and moreover, indicate that two velocity components from different structures are required. With higher resolution observations these structures could possibly be dis-entangled. Taking into consideration the luminosity of G17.64 (1 × 105L ) and that it appears as the main dominant and
iso-lated source in continuum emission, a mass in the upper limit of the matching model range ∼20 M is more probable.
4.2. Possible origins of the SiO emission and kinematics The presence of SiO means silicon must have been released into the gas phase by sputtering, likely due to a C-shock (e.g.Schilke et al. 1997; Gusdorf et al. 2008) or by grain–grain collisions, shattering or vapourising the grains (Jones et al. 1996;Guillet et al. 2011). The maximal velocities of the SiO we detect (relative to the VLSR) are ∼20 km s−1marginally slower than the typical
C-shock speeds >25 km s−1(Caselli et al. 1997). Observationally
our SiO detection is only resolved by approximately three to four beams and it is plausible that the even higher velocity emission
